Collaborative Research: New Standard Stellar Population...

Project Description

Collaborative Research: New Standard Stellar Population Models

Few classes of models are more ubiquitously spread throughout the field of astronomy than stellar population models. They are used any time that the luminosities and colors of galaxies or individual stars are important, either as ends in themselves or as a complication to be subtracted in order to arrive at the desired science. This document outlines (1) an overview of the field-spanning impact of such models, (2) limitations in the models presently available, and (3) our proposal for a greatly improved, maximally flexible array of models that will output a comprehensive array of colors, magnitudes, SBF magnitudes, line strengths, and fluxes for integrated-light applications, as well as multicolor isochrones and Hess diagrams for those interested in individual stars. One of the most novel features of the new models is that they can be interpolated for any mixture of H, He, C, N, and O, plus three options for different patterns of alpha-element enhancement. This feature should allow, for the first time, reliable mean age estimates for giant elliptical galaxies that are known to have non-solar element mixtures, with a realistic goal of 10% age uncertainty. Chemical evolution studies of all kinds will benefit enormously from these models. The final models will be interfaced with web front-ends so that interpolated models can be generated on the fly, with cosmological fading, redshifting, and evolutionary corrections. (4) Scientific goals for elliptical and lenticular galaxy age and abundances are described next: ages for nearby galaxies and clusters with high quality data and an investigation into the UV flux of elliptical galaxies. (5) The final main section of this proposal will give a specific implementation plan and a timeline for completion. (6) Appendices on model atmospheres and model error propagation are provided at the end.

1. Stellar Population Models Permeate Astrophysics

Stellar population models use stellar evolutionary isochrones plus stellar flux information summed over a single-burst stellar population to give integrated-light observables like colors and spectral line strengths. They have a variety of ingredients, but always two main components: stellar evolution to populate the stars in the log L, log T diagram, and stellar observable properties (colors, etc.) so that the theory can be related to observation. We make little distinction between isochrones and integrated-light models since the latter are the former under an integral sign. To condense the “modern era” of stellar populations to few sentences, Bruzual (1981) showed how stellar evolution could be used to predict the colors and luminosities of high-redshift galaxies. At about the same time, O’Connell (1980) showed that age information could be extracted from a single integrated light spectrum, but only if the metallicity was

Worthey et al. - 1 - Project Description

known in advance. Improvements in stellar evolution theory (e.g. VandenBerg & Bell 1985) bolstered hopes that precision age information could be extracted from cluster color-magnitude diagrams, and allowed the addition of metal abundance as a parameter in integrated light models. With metallicity added and with the addition of spectral absorption feature indices (Worthey 1994) a method was developed to extract a mean age and metallicity simultaneously from Balmer and metallic absorption features measured in a single integrated light spectrum of objects with ages between ~0.5 and 14 Gyr. Leitherer (1995a,b) pioneered work on young populations, where the main degeneracy is between age and slope of the initial mass function (IMF), and Leonardi & Rose (1996) found a way to disentangle burst age from burst mass in post-starburst objects. Today, with many isochrone grids available, many integrated light models have been made, some generalized, some specialized. For what are they used? The following incomplete list gives a flavor for the multiplicity of applications that could be addressed by better, more flexible models such as the ones we propose. We’ve suppressed references for brevity.

• Open and globular cluster ages from color-magnitude diagram studies. We expect that the chemical flexibility of our isochrones will provide a great boon to this endeavor.

• Star formation histories of local dwarf galaxies are derived from color-magnitude diagram decomposition, including the chemical history in favorable cases.

• Star formation history of the Milky Way through star counting programs. • Extragalactic globular cluster systems within ~10,000 km/s have color

distributions that can be statistically related to galaxy formation scenarios, and cluster systems within ~1500 km/s can be studied spectroscopically to get individual ages and metal abundances for each.

• Quiescent (E, S0, spiral bulges) galaxies are amenable to mean age determination via stellar population models, and models with the proposed flexible chemistry are crucial to making progress in this area.

• Light/heavy element ratios, especially in elliptical galaxies, are known to vary from the solar ratios. Measurement of abundance ratios is much more accurate when the models can be adjusted for different chemistries.

• Surface-brightness fluctuation magnitudes remain an exciting distance measurement tool. Stellar population models give a-priori values that skip almost all distance-ladder steps. The same models are the only way to estimate k-corrections for SBF.

• Post-starburst objects give crucial insight into galaxy cluster formation and rely heavily on stellar population models.

• Starburst objects are studied with the aid of stellar population models. • Quasars and AGN occasionally have to have the starlight subtracted in order to

reveal more about the nebular diagnostics and synchrotron continuum. Conversely, the hosts of these powerful objects can be studied through analysis of their starlight using stellar population models.

• Cosmological studies of galaxies utilize stellar population models to account for stellar population evolution and k-corrections as a function of lookback time.


• Spectroscopic studies of individual galaxies to z>1 are now possible with large telescopes, with abundance and age measurements possible through stellar population models in cases of sufficient signal.

• Chemo-dynamical galaxy evolution models (N-body, hydro, SPH) are coming to rely more and more on stellar population models so that predictions can be compared with observations.

• Chemical evolution, typified by the G-dwarf problem, can be extended to external galaxies as far as the Virgo cluster via color-magnitude diagrams of the red giants, interpreted through stellar population models.

This list illustrates the broad impact of stellar population models. We now turn to the question of what can be improved in current models to aid the variety of endeavors in which they are used.

2. Fix-it List: Flaws and Omissions in Today’s Models

There are many stellar population models available today (e.g. Blakeslee et al. 2001; Bressan et al. 1994; Bruzual & Charlot 1993; Buzzoni 1993; Dorman et al. 1993; Kodama & Arimoto 1997; Liu et al. 2000; Salasnich et al. 2000; Schulz et al. 2002; Schiavon et al. 2002a,b; Tantalo et al. 1998; Vazdekis et al. 1997; Worthey 1994). Unfortunately, to use such a model is to enter a confusing world of dubious quality control and conflicting results. These days, with both isochrone sets and synthetic fluxes now electronically available, it is a matter of a few days of programming to unite the two to produce a model for integrated light. The authors are of the opinion that the stellar population field has gotten a reputation for crazy results because of the proliferation of hastily constructed models plus misuse of properly constructed ones. The main causes for discrepancies are the stellar evolution ingredients (Worthey 1994; Charlot et al. 1996; Yi 2003). That is, the number, temperatures, and luminosities of the stars are a bigger uncertainty than the observational character of the stars themselves. We give five tutorial examples of where models deviate from reality, plus notes on how we will improve the situation. Example 1: K-band Surface Brightness Fluctuation Magnitudes. One of the persistent preconceptions about stellar population models for integrated light is that, because they sum over all stars, errors in evolution or stellar fluxes will be muted, or even “averaged away” if the error sources are random. In reality, errors in the input ingredients are as often boosted as they are damped. One particularly vulnerable part of the H-R diagram is the red tip of the red giant branch (RGB) and asymptotic giant branch (AGB) where stars are cool, luminous, and large. Present-day model-builders have a spotty record for adjusting the temperatures of these giant branches to match metal-rich populations, as illustrated in Figure 1 for populations of solar metallicity. The figure shows surface brightness fluctuation (SBF) magnitudes in the K band. SBF magnitudes are a useful distance indicator for local galaxies, and depend on the luminosity function of the brightest stars present in the stellar population. In the K


band, the coolest, most luminous giants dominate the corresponding SBF magnitude (for an ordinary magnitude, the K light is more equally shared among all locations on the RGB and AGB branches). Figure 1 shows a variety of zeropoints and population age sensitivities among different authors. This scatter nearly vanishes if the theoretical giant branches are forced to pass through those of calibrating open clusters such as 47 Tuc, M67, and NGC 6791 at accepted age and metallicity values. The proposed new models will match calibrator cluster temperatures (and significant new work will be done on temperature calibration; c.f. section 4). Figure 1. K-band SBF magnitudes from several stellar population models (Liu et al. 2000, Blakeslee et al. 2001, Buzzoni 1993, and Worthey 1994). The large volatility is due primarily to temperature drifts among bright giants, secondarily to the treatment of the luminosity function of RGB tip and AGB stars, and to a small extent to the different stellar color-temperature relations. The desired accuracy is 0.1 mag for 5% distances.

All the models in Fig. 1 except W94-with-Padova-isochrones were specifically built to compute SBF magnitudes, but scatter by more than a full magnitude. In V- or I-band, the scatter between authors is reduced because RGB-tip and AGB stars are fainter in bluer passbands. Example 2: The mean age of 53w091. This example shows both a scatter of results from different model authors and also the difficulty of finding an age for a galaxy due to the “age-metallicity degeneracy” where an age change has nearly the same effect as a metallicity change on the integrated-light spectrum. A z=1.55 galaxy (LBDS 53W091) was found to have the red absorption-line spectrum of a passively evolving stellar population (Spinrad et al. 1997). The spectrum is rest-frame ultraviolet, the region of the spectrum where fluxes become uncertain to the point where flux errors may compete with stellar evolution errors.


Figure 2. The spectrum of galaxy 53W091 (the ratty line) along with 3 models from Worthey (1994) that illustrate the age-metal degeneracy. Unless one knows the metallicity to better than 0.2 dex, the age is unknown to a factor of two.

The age range covered by integrated flux models by various authors, all of which match the UV spectrum of this galaxy fairly well, ranged from less than 1 Gyr to more than 12 Gyr. The model that we could check all had similar main sequence turnoff (MSTO) temperatures, so the main problem in this example is the UV stellar flux library. This problem can be almost 100% fixed with suitable comparisons to new HST spectra or even old IUE spectra (We plan just this. See section 4 and appendix 2). The additional lesson that we learn from this galaxy is that simply matching a swath of spectrum with a model does not lead to reliable science because of the age-metallicity degeneracy effect. As shown in Figure 2, models of different ages fit the spectrum equally well if one is allowed to change the metal abundance slightly. Example 3: The 47-Tuc age discrepancy. One important application of stellar population models is to estimate the mean age and metal abundance of a stellar system by looking at particular spectral features in its integrated spectrum. Begun with Balmer features arrayed against metallic features (W94) the search continues for the best spectral features to apply to this problem. Jones & Worthey (1995) introduced an Hγ index that appeared to have extraordinary ability to measure the mean age of old stellar populations. Sensibly, this index was tested on globular cluster 47 Tucanae, for which both an integrated spectrum and a stellar main sequence turnoff age can be obtained. Gibson et al. (1999) found that the spectroscopic age exceeded 20 Gyr in all of their models, while the accepted age from color-magnitude diagram (CMD) studies is around 11 Gyr. This result was viewed with some combination of amusement, alarm, and embarrassment, and illustrates that (1) we are finding very effective, age-sensitive spectral indices, and (2) modeling them incorrectly.


In the case of the 47-Tuc age discrepancy, the answer has apparently been found by Schiavon et al. (2002a,b), and boils down to a small collection of effects. (1) Inclusion of alpha-element enhancement and heavy-element diffusion in the isochrones both reduce the discrepancy via modest modifications of stellar temperatures and lifetimes, and (2) the red giants need careful attention, especially as regards their luminosity function. This is true even though Hγ is strong only in warm stars near the main sequence turnoff, and therefore, one would think, insensitive to changes in red giant spectra. Example 4: The age – chemical mixture interdependency. The phrase “3/2 age-metal degeneracy” (W94) refers to the fact that a 20% increase of metal abundance has almost exactly the same effect on the integrated spectrum of a stellar population as a 30% increase in age. Use of carefully chosen spectral features (see last example) can defeat this degeneracy to simultaneously measure a mean age and a mean metallicity from one integrated-light spectrum. If 10% ages are our goal, however, many effects buried in the noise before reveal themselves as crucial. These include horizontal branch morphology and treatment of blue stragglers (more on these below), but also heavy element mixture (mixture, not amount). Heavy-element mixture can change the shape of the underlying isochrone, leading to incorrect age derivations. An example of this is shown in Fig. 3, where an increase in relative oxygen abundance while [M/H] stays constant drives an age error of 70%. Abundance ratio changes affect stellar temperature by altering the balance between various sources of opacity, and thus changing the stellar radius and effective temperature. Figure 3. Isochrones showing the effect of increased oxygen abundance. In this example, adding oxygen lowers the turnoff temperature by 150 K without affecting the giants, leading to an age estimate that is 70% too old. This example is extreme, but it underscores the need for isochrones with flexible chemical mixtures if accurate ages are desired. One important note about Fig. 3 is that it was computed at low metallicity for application to globular clusters. Quite a few “alpha enhanced” isochrones are available at the metal-poor end but similarly detailed work for metal-rich populations such as comprise the bulk of the stars in the universe has not been carried out. A few authors (Weiss et al. 1995; Salaris, unpublished, Thomas et al. 2002; Salasnich et al. 2000) have made forays in this direction, concentrating primarily on [Mg/Fe]. Additionally, and importantly, nobody so far has coupled appropriate stellar atmospheres (that share the same altered abundances) to the interior models. Atmospheres have a huge effect on stellar radii and effective temperatures, and the full implications of this absolutely must be explored as a function


of abundance mixture. We propose to make fully consistent interior-plus-atmosphere stellar models. The effects of chemical mixture on age can be handled by using models with the correct chemical mixture to derive the age. We adopt as our benchmark models good to 10% in absolute age (i.e. about the same as well-studied star clusters) and 5% relative age as a function of chemical mixture. We plan to achieve this relative accuracy with our proposed models. Example 5: We can estimate but not measure individual element overabundances in elliptical galaxy spectra. Figure 4a marks some spectral features due to individual elements or simple blends that are visible in the low-resolution spectra of bulges, elliptical galaxies, and lenticular galaxies, and Figure 4b shows models and observations for the case of Mg versus Fe indices. The models are of various ages and abundances (scaled-solar, of course, no mixture variations are possible in these old models). Large elliptical galaxies, not having a scaled-solar abundance mixture deviate fairly strongly from the model line. These galaxies look like they lie about 0.3 dex too high in Mg relative to Fe (Worthey et al. 1992; Weiss et al. 1995; Trager et al. 2000a, b). Note that the observational error bars compared to the throw of the models indicates an abundance precision of 0.03 dex (small error bar) to 0.06 dex (large error bar), indicating highly accurate abundance measurements are possible with today’s data provided the models can deliver such accuracy, which, today, they cannot. The “measured” [Mg/Fe] overabundance in Fig. 4 is true in the limit where the isochrone temperature structure does not change with abundance mixture, but as we argue in Example 4, temperatures can change fairly strongly. It is clear that we need models that match real galaxies in detailed composition if we are to convert these estimates to real measurements.

Figure 4. Panel a) shows a galaxy spectrum with prominent spectral features (blends) labeled by the major element or molecule that contributes. This is a favorable case, with low velocity dispersion and therefore crisper spectral resolution. Panel b) illustrates how age and metallicity effects, shown by model lines with scaled-solar [Fe/H] indicated, lie nearly orthogonal to the effect of Mg enrichment in a sample of elliptical galaxies (triangles).


Many elements, like Mg, Ca, N, Fe, and Na are readily measurable in the spectra of quiescent galaxies (Fig. 4). With some work, we expect several others to be attainable with comparable precision, notably Si, Ti, C, Sc, Ti, and V (Worthey 1998). Oxygen and helium, however, have no absorption features that are easily measurable in low-resolution spectra, so we must rely on secondary effects to infer the abundance of these elements. Tripicco & Bell (1995) and Houdashelt et al. (2003) computed synthetic spectra from which Lick indices are measured for 3 representative stars as a function of mixture. Inspection of these spectra give strong hints that it may be possible to construct indices that measure O, Ni, Cr, and a few other elements even where no obvious spectral absorption bump exists. Clearly, an entire grid of synthetic spectra with chemical mixture as an additional variable could be combined with isochrones to fully explore this problem. Exactly such a grid is proposed here.

3. Our Proposal: New Standard Stellar Population Models

We give here the parameters that shape the proposed effort.

o Flexible chemistry will be the hallmark of these models. Abundance will be characterized by separate parameters for H, He, C, N, O, and three choices for “alpha”/heavy or light/heavy element mixtures. This has never been done in anything approaching a self-consistent way, and we feel that the ultimate scientific dividends in the areas of chemical evolution and precision age estimates will exceed what we can predict at this time.

o Theoretical atmospheres, opacities, and fluxes will be calculated using the PHOENIX code (Hauschildt & Baron 1999). We think that this is the most accurate code available, with non-LTE effects and spherical atmosphere geometry included. The spectra will be used to calculate synthetic isochrone colors and summed to produce integrated light spectra. The opacities will be used in the DSEP stellar evolution code. (See appendix I for details about PHOENIX)

o Interior opacities are now available for arbitrary chemical mixtures from LLNL, and they will be computed to match our abundance scheme.

o Stellar evolutionary isochrones will be calculated for stellar population ages older than a few million years with a sophisticated evolution code. Chaboyer's stellar evolution code, DSEP, is a descendant of the Yale code (Guenther et al. 1992). This code accurately models the evolution of stars in the mass range 0.5 to 5 solar masses. DSEP is able to follow the evolution of stars starting on the pre-main sequence and ending at the tip of the asymptotic giant branch. DSEP uses state of the art input physics, including the OPAL high temperature opacities (Iglesias & Rogers 1996), nuclear reaction rate data from Adelberger et al. (1998) and an equation of state which includes the Debye-Hückel correction (Guenther et al. 1992). This equation of state yields stellar models which are in good agreement with those derived using the OPAL equation of state (Rogers 1994; Chaboyer & Kim 1995). Helium and heavy element diffusion are including using the Thoul et al. (1994) diffusion coefficients. The diffusion near the surface is inhibited, so that the models are able to accurately reproduce observations of Li and other heavy elements in Pop II stars, while at the same time yielding solar models that are in excellent agreement with helioseismology (Chaboyer et al. 2001). Convective core overshooting is


implemented in the manner described by Woo & Demarque (2001), which follows Roxburgh (1992) in implementing an upper limit to overshoot. This technique has been shown to yield stellar evolution models that are in good agreement with a variety of open clusters color-magnitude diagrams. DSEP will be modified for this project so that it will use low temperature opacities and model atmospheres (for surface boundary conditions) generated by PHOENIX.

Figure 5: The best observed bimodal, metal-rich horizontal branch (top: from open cluster NGC 6791, membership data from K. Cudworth, private communication) compared with two representative synthetic horizontal branches (middle and bottom). The synthetic HBs were computed for uniform mass-loss: any mass-loss between zero and MC is equally probable (and this clearly cannot match the observations, so a bimodal mass-loss mechanism is indicated).

o For old populations, synthetic horizontal and asymptotic giant branches will be

constructed from sets of late-stage evolutionary tracks in which the envelope mass is allowed to vary. The synthetic horizontal branches in Fig. 5 were built by us, with Dorman (1993) tracks, illustrating the problem of bimodal HB temperature distributions in metal-rich populations. The hot clump of HB stars are responsible for the excess UV light in elliptical galaxies (Ferguson 1999), but bimodality does not arise naturally from the stellar lifetimes of the stars present: the clump of hot stars have had 0.3 to 0.4 solar masses of material stripped from their outer envelopes during their tenure on the RGB, almost down to the helium core. Our synthetic HB and AGB distributions will be parameterized and tailored to match observations (see appendix 2), by predictions from binary evolution theory, or, if that does not yield predictive results, by an empirical mass-loss that matches HB morphology in all calibrating star clusters. This investigation will be carried out at Dartmouth. See further comments in Section 4.

o A comprehensive array of model outputs will be presented. Photometry will be given in the Johnson/Cousins, HST, Strömgren, Washington, and SDSS systems. Surface brightness fluctuation (SBF) magnitudes will be given in the same passbands. Line strength indices on the Lick/IDS system and extended Rose system will be given. Synthetic spectra at point spacing <1 Å in the optical will be given, along with spectra based on empirical stellar libraries (i.e. the spectra will be empirical at scaled-solar abundances but will step away differentially as the element mixture changes).


o Results from the models will be made available to the community through the data centers and real-time interpolations will be available over the web. Like Worthey’s existing dial-a-galaxy page (http://astro.wsu.edu/worthey/dial/dial_a_model.html) or as promised soon from the IAP group (http://www.iap.fr/pegase/), we will present dial-an-isochrone, dial-a-star-color, dial-a-galaxy, and dial-a-cosmology forms, along with access to data tables and interpolation routines.

Here are the items relating to stellar evolution or synthetic spectra for which empirical corrections or approximations may need to be applied. Too much empiricism may limit the predictive power of the final models, so we wish to minimize it as much as possible. On the other hand the models must match all calibration tests, and if they don’t do so “naturally,” compliance will be forced upon them.

• Stellar evolution near the tip of the first-ascent RGB is tricky because computed effective temperatures are seen to drift from observed temperatures for metal-rich giant branches due to the difficulty of theoretically matching the bloated and vacuous atmospheres of these stars. Temperature systematics will be all but eliminated by comparison with our cluster network (appendix 2). If DSEP plus opacity from PHOENIX cannot match Teff naturally, we will apply empirical shifts.

• Synthetic spectra, although rapidly improving, still cannot match an empirical spectrum at the 1% level at, say, 5 Å resolution. We will follow two schemes to improve the situation to maximize the potential for age determination and chemical analysis of galaxies. First, as has now become standard, line-strength indices will be attached, where the indices come from empirical spectra. Changes due to changing abundance pattern will be computed differentially from the synthetic spectra. Second, several stellar spectral libraries now exist of sufficient quality to attempt a differential correction for whole swaths of spectrum at once. Improvement in spectral modeling is crucial for galaxy ages and for chemical evolution studies.

• Stellar temperature conversion to colors and fluxes is usually not the dominant source of error for integrated light models, but can affect stellar CMD studies seriously. Errors for colors will be propagated through to model outputs (see appendix 2). We expect synthetic colors to be good to 0.01 – 0.02 mag without any fudging.

• The treatment of mass-loss on the RGB and AGB and for supergiants is tricky because of the inherently episodic nature of mass loss (e.g. Origlia et al. 2002) plus the apparent two-mode mass loss needed for hot clumps of HB stars (observed clusters, elliptical galaxy UV bumps). We will tune our mass-loss parameterization to match the observed morphologies of horizontal branches at all metallicities. Calibrating clusters are unfortunately sparse. NGC 6791 at [Fe/H] = 0.3, a handful of solar-metallicity cluster stars, and globular clusters comprise the complete workable data set. We will explore the hypothesis that binarism is responsible for the bimodality of metal-rich cluster horizontal branches, but if a “natural” mechanism is unconvincing we will fall back on the two-Gaussian mass loss approach taken by most authors today. The latter is clearly second-best because we will not be able to fully predict the impact of abundance ratio changes on horizontal branch morphology, but this will not seriously impact our 10% age goal or our goal of measuring abundance ratios because blue HB light is an extremely minor player in the visible spectral region.

• We will include tracks for carbon stars when they appear in the population and blue stragglers under the hypothesis that they are merged binaries.


http://astro.wsu.edu/worthey/dial/dial_a_model.html

4. Science Goals

In order to endure the painstaking task of amassing the largest collection of stellar fluxes and evolutionary isochrones ever assembled and patiently checking the results against spectrophotometry, photometry, clusters, and other data, accompanied by a lot of input/output programming and web page construction, we reward ourselves at the end with a focused collection of science goals. Although considerable “broad impact” will come from the existence of the models themselves, we expect that these scientific results will also be influential in refining our view of galaxy evolution.

First, we want to measure 10% absolute ages from a single spectrum of a simple (or near-simple) stellar population. We emphasize that current age uncertainty is almost entirely a model issue (due to the flaws and omissions outlined in Section 2). Integrated-light index data of high enough quality for 10% ages exists today in small amounts for some globular clusters and many bulges, elliptical galaxies, and lenticular galaxies, and we expect more soon from either other workers or our own partially reduced data on a volume-limited sample of nearby galaxies. We will apply our new models to the data by dialing them to the right element mixture, then using the Balmer-metal relations to derive an age. We can then correctly assess errors (for the first time) by dialing up other abundance mixtures that are within the errors and assessing age scatter in multiple Monte Carlo realizations. Second, we will measure [X/Fe] element ratios for C, N, Ca, Si, Sc, V, Ti, and Mg for the same set of cluster and galaxy data, also measuring He, O, Ni, and Cr if it turns out to be possible. As for ages, observational spectra of sufficient quality already exist (c.f. Section 2, Example 5: 0.03 dex precision); to measure the abundance ratios is solely a matter of producing flexible enough models. Since we will output synthetic spectra of flexible element mixture we will know at the end of the modeling effort whether or not modest element abundance changes are reflected in measurable spectral changes. We will check the rest of the periodic table of elements to see if there are other measurables (K? Co? Cu? Sr?) but it is hard to predict at this point how much success we should expect. Third, we will make an attempt to discover the origin of the scatter in UV-visible color of elliptical galaxies. Background: elliptical galaxies possess an obvious “bump” of UV light, peaking at roughly 1500 Å (Burstein et al.1988). When the bump is strong it is dominated by “Extreme Horizontal Branch” (EHB) stars such as those illustrated in Fig. 5 (Ferguson 1999; Dorman et al. 1995). There is about a factor of ten spread in 1550-V color among elliptical galaxies. The obvious parameters of overall metal abundance, age (as best as we can measure it today), luminosity, or velocity dispersion cannot amongst themselves predict 1550-V, leaving much to be explained by an additional parameter. We will use our new measurables ([X/Fe], higher-precision age) plus our flexible HB/AGBs to make a solid try at discovering the missing parameter. Even in the event that we cannot find a “natural” bimodal mass-loss solution, we expect to make considerable progress on this question.


5. Implementation Plan and Timeline This collaborative project pools a team of three PIs, experts in stellar evolution (Dartmouth), stellar atmospheres (Oklahoma), and population synthesis (Washington State). It involves substantial amounts of computer time, including supercomputer time for the PHOENIX runs, but it also involves a heavy load of hands-on work so that we expect the project to be manpower-limited. We therefore request a postdoc at WSU, graduate students at OU and Dartmouth, summer stipends for Baron and Worthey, and an undergraduate researcher at WSU. E. Baron at OU and his student will oversee the PHOENIX runs and do the necessary validations and checks on stellar properties. B. Chaboyer and his student will concentrate on post-helium-flash evolution, including HB morphology, mass-loss, and binarism.The WSU postdoc will be trained by Chaboyer at Dartmouth but do stellar evolutionary calculations at WSU. The program in astronomy at WSU currently supports no graduate students, so a postdoc is a necessity. The postdoc will perform the stellar evolution grid calculations, primarily from main sequence through helium flash (while later stages will be calculated at Dartmouth), and, in year 3, help with quality control and synthesis. The undergraduate at WSU will in the first year work on assembling the latest stellar photometry and temperature data. In years two and three the undergraduate will work on data cross-checking and the public web interface. Trager and Worthey will be in charge of overall quality assurance, and Worthey will supervise the postdoc and undergraduate student. Year 1 will see the calculation of atmosphere fluxes, Rosseland mean opacities, the final distillation of observational constraints, and the beginning of the stellar evolution calculations. The first year will also include the collection of the best (star) cluster observational data, spectrophometry (primarily from an HST snapshot program called NGSL, PI M. Gregg, and the F-star spectrophotometry HST Legacy program, PI R. Peterson), and additional data on the stellar temperature scale including additional field star photometry. This will be packaged into an empirical color-temperature grid and published. The synthetic fluxes will be evaluated against photometric and spectral observations. Year 2 will see the completion of stellar evolution calculations through the Helium flash, and the bulk of the calculations of the synthetic HB and AGB. This phase will require final decisions regarding the adoption of a first-principles mass-loss scheme or a more empirical scheme to constrain the HB to conform to observations. The population synthesis and isochrone interpolation routines, already prototyped in Worthey (1994), will be expanded. Empirical spectral libraries will be added as an option to the synthetic fluxes. Year 3 will see the most aggressive quality control activity. The model outputs will take their final form. The web-based user interfaces will mature. We will address the scientific questions of 10% ages from integrated light and whether it is possible to measure the abundance of oxygen and helium from second-order spectral effects. The models will be announced to the community.


Appendix Ι: PHOENIX

PHOENIX is a general, fully relativistic, non-local thermodynamic equilibrium (NLTE) stellar atmosphere and spectrum synthesis code. Peter Hauschildt (UGA), Baron, France Allard (Lyon) and collaborators have further developed PHOENIX for modeling hot and cool stars, brown dwarfs, giant planets, and supernovae. In order to meet the large CPU and storage requirements of PHOENIX, we have adapted the code to run effectively on parallel machines in both shared and distributed memory modes. PHOENIX treats the interaction between radiation and matter in fully detailed NLTE. It uses state of the art radiative transfer methods for relativistically expanding atmospheres and it includes extremely complex model atoms for iron-peak elements. It treats deposition of radioactive energy by gamma rays and fast electrons, and it can handle various boundary conditions including external illumination as well as nebular conditions. In the current version (13.1) of PHOENIX the fully relativistic radiative transport equation is solved for continuum and line radiation simultaneously and self-consistently using an operator splitting technique (see http://phoenix.physast.uga.edu). Such a method also is used to solve the full multilevel NLTE transfer and rate equations for a large number of atomic species with more than 11,000 energy levels and 100,000 primary NLTE lines, including non-thermal processes. Radiative equilibrium is enforced to ensure global energy conservation in the co-moving frame, with line blanketing and non-thermal processes taken into account. This condition is a general condition of energy conservation, and is suited to coupling our calculations with the results from hydrodynamic codes.

Figure 6. Ultraviolet spectrum of Deneb from the IUE satellite compared to LTE (local thermodynamic equilibrium) and non-LTE PHOENIX spectra. Ultraviolet spectral synthesis is challenging because of non-LTE effects and also because of inaccurate line lists. The PHOENIX spectra match real stars much better than LTE models such as the Kurucz models that are the staple for current stellar population models.


In a detailed radiation transport calculation the goal is to keep track of the populations of each state in each model atom, the distribution function of the radiation field, and the interactions between them. The enormous number of transitions that have to be included in realistic models of atmospheres with large scale velocity fields requires an efficient method for the numerical solution of the multilevel NLTE radiative transfer and model calculation problem. Simple approximations such as the Sobolev method are not accurate when lines overlap strongly and form a significant pseudo-continuous opacity. Classical techniques such as complete linearization or the equivalent two level atom become computationally prohibitive. Anderson’s (1989) multi-group scheme or extensions of the opacity distribution function method (Hubeny & Lanz 1995) cannot be used for moving media such as supernovae. Methods that are based on partial linearization schemes tend to be numerically unstable, and frequently encounter convergence problems because of the highly nonlinear and nonlocal couplings that dominate supernova atmospheres. Special methods that allow us to address all these effects without further approximations have been developed and thoroughly tested (Hauschildt & Baron 1999). As more transitions are included the number of wavelength points that must be included also increases, since each line profile must be resolved (although we have lessened this requirement by using direct opacity sampling). Therefore we have developed an algorithm to parallelize over wavelengths that allows us to take advantage of massively parallel processors that are available at national supercomputer centers such as NSF NPACI centers. Although code development is an ongoing process, the parallelized version of PHOENIX has reached full production status after an intense testing period and parallel supercomputers are now used on a regular basis for practically all PHOENIX calculations.

Appendix II. Calibration Procedures

To maximize the usefulness of the new stellar population models it is highly desirable to attach a “model error” to each predicted observable. In this way, users of the models will have a good idea of how far they can push model interpretations. Understandably, given how much work is involved, few of the current stellar population models make any attempt to calculate model errors. Those that do limit themselves to post-hoc comparison in which models from different authors are overplotted, as exemplified in Charlot et al. (1996). The only attempt to calibrate model errors via propagation of errors from the input ingredients that we could find was Worthey (1994; W94), who was able to estimate errors in colors quite well, but errors in spectral indices quite poorly. Here is a brief list of procedures that we will use in verifying the quality of the new models. Model fluxes: Synthetic fluxes will be compared to observed spectrophotometry, especially the new NGSL library from an HST/STIS snapshot survey (PI: M. Gregg, co-I Worthey). These comparisons will quantify how well synthetic absorption features match empirical measurements. The synthetic fluxes are also used to compute magnitudes and colors multiplying the filter transmission function, (passband) Pλ with the flux Fλ and integrating over wavelength: Mag = -2.5 log ∫FλPλdλ + ZP, where ZP is a zeropoint correction to shift the result to the stellar magnitude scale. The zeropoint is most often calculated by matching a model to the spectrum of Vega. Colors are matched to Vega, and the overall bolometric zeropoint is chosen to match the Sun. Both Baron and Worthey (1994) have well-tested synthetic photometry codes already operational. Once we have zeroed the scale for magnitudes and colors we compare with established


bolometric correction sequences as a function of temperature to assess error. Color-color relations will be matched for all photometric systems to obtain an uncertainty for each individual color. This was one of the most useful comparisons in W94, and revealed a variety of systematic color drifts. Finally, color-temperature relations will be examined, with special attention to V-K and V-I because of their insensitivity to metal abundance variation (von Braun et al. 1998). The net result will be an accounting of flux errors. The flux uncertainties can then be folded in with uncertainties in the stellar evolution prescription. Stellar evolution uncertainties tend to dominate in integrated light, but CMD (color-magnitude) studies will be most affected by color uncertainties. The key to minimizing errors in the isochrones is to establish a network of observed clusters that cover the available parameter space of age and metal abundance. A sequence of globular clusters will be used, plus several key open clusters including the Pleiades, the Hyades, Messier 67, NGC 6791, Melotte 66, NGC 2477, NGC 2204, and old, metal-rich globular clusters near the Galactic bulge. Worthey is already involved in investigating a superset of this cluster list in connection with his Space Interferometry Mission duties. The WIYN Open Cluster Project has promised early access to their data for overlapping objects (see attached letter). The CMDs of these clusters will be reduced to fiducials in log L, log Teff space, and the isochrones will be directly compared at various slices in log L. The temperature uncertainties will be known from the previous work with fluxes. Some additional error will creep in from uncertain distance moduli, reddenings, ages, and metal abundances for the clusters, but (a) we are after an ensemble average, and (b) we are looking for key diagnostics: relative temperature of main sequence turnoff, temperature of giant branch stars, luminosity of the horizontal branch, and luminosity function, and (c) the clusters are chosen for low reddening and how well-studied they are. We can do much better than in 1994. In particular, JHK colors from 2MASS will homogenize and solidify the temperatures of cool stars. For our synthetic HB/ABG effort, this data plus more will be utilized. Such data is incomplete only at the metal-rich end. With these efforts, uncertainties can be estimated with fairly good confidence, but we will confirm the estimates of uncertainty with comparisons to empirical data after the isochrones and fluxes have been combined. Various galaxy and integrated-light star cluster color-color, index-index, and SBF (surface brightness fluctuation; e.g. Tonry et al. 2001) magnitude diagrams will reveal if the calculated uncertainties are realistic. The network of clusters mentioned above will be examined in observed color-magnitude space to confirm that the temperature transformations and bolometric corrections operate as expected. The V-I, R-I diagram for SBF magnitudes in early-type galaxies is an extremely revealing diagnostic: no models have yet reproduced this diagram in a satisfactory way. Finally, Lick index fitting functions and new designer indices tailored to measure individual element ratios will be much better determined than previously from the stellar sample because of the large number of high-quality spectra of stars of disparate atmospheric parameters now available. For the first time, realistic errors can be estimated by subsampling within the complete stellar spectral catalog. If the uncertainties are calibrated, we feel that the final result will be demonstrably reliable and thus a useful standard for years to come.


Collaborative Research: New Standard Stellar Population...

Documents

Transcript of Collaborative Research: New Standard Stellar Population...